ASTRONOMY & ASTROPHYSICS
NOVEMBER I 2000, PAGE 437
SUPPLEMENT SERIES
Astron. Astrophys. Suppl. Ser. 146, 437–464 (2000)
Dust features in the 10-µm infrared spectra of oxygen-rich evolved
stars
A.K. Speck1 , M.J. Barlow2 , R.J. Sylvester2 , and A.M. Hofmeister3
1
2
3
Astronomy Department, University of Illinois at Urbana/Champaign, 1002 W. Green Street, Urbana, IL 61801, U.S.A.
e-mail: [email protected]
Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK
Department of Earth & Planetary Science, Washington University, St Louis, MO 63130, U.S.A.
Received December 22, 1999; accepted August 15, 2000
Abstract. We have analyzed the 8 − 13.5 µm UKIRT
CGS3 spectra of 142 M-type stars including 80 oxygenrich AGB stars and 62 red supergiants, with a view to
understanding the differences and similarities between the
dust features of these stars. We have classified the spectra
into groups according to the observed appearance of the
infrared features. In each case the normalized continuumsubtracted spectrum has been compared to those of the
other stars to find similarities and form groups. The dust
features of the AGB stars are classified into six groups:
broad AGB, where the feature extends from 8 µm to about
12.5 µm with little structure; broad+sil AGB, which consists of a broad feature with an emerging 9.7 µm silicate
bump; and four silicate AGB groups in which a “classic” 9.7 µm silicate feature gets progressively narrower.
Likewise, the supergiant spectra have also been classified
into groups, however these do not all coincide with the
AGB star groups. In the supergiant case we again have
six groups: featureless, where there is little or no emission
above the continuum; broad Super, where the feature extends from about 9 µm to about 13 µm; and four silicate
Super groups, which again show a progression towards the
narrowest “classic” 9.7 µm silicate feature. We compare
the mean spectrum for each group, which yields two main
results. Firstly, while the “classic” silicate feature is essentially identical for both AGB stars and red supergiants,
the broad features observed for these two stellar types are
quite different. We suggest that the dust in these two environments follows different evolutionary paths, with the
dust around Mira stars, whose broad feature spectra can
be fit by a combination of alumina (Al2 O3 ) and magnesium silicate, progressing from this composition to dust
dominated by magnesium silicate only, while the dust
around supergiants, whose broad feature can be fit by a
combination of Ca-Al-rich silicate and Al2 O3 , progresses
Send offprint requests to: A.K. Speck
from this initial composition to one eventually also dominated by magnesium silicate. The reason for the difference in the respective broad features is not clear as yet,
but could be influenced by lower C/O ratios and chromospheric UV radiation fields in supergiant outflow environments. The second result concerns the 12.5 − 13.0 µm
feature discovered in IRAS LRS spectra and widely attributed to Al2 O3 . This feature is seen predominantly in
the spectra of semiregular variables, sometime in Miras
and only once (so far) in supergiant spectra. We argue
that it is unlikely that this feature is due to Al2 O3 or, as
has more recently been suggested, spinel (MgAl2 O4 ), but
could be associated with silicon dioxide or highly polymerized silicates (not pyroxenes or olivines).
Key words: (stars:) circumstellar matter — stars:
mass-loss — stars: AGB and post-AGB — (stars:)
supergiants — infrared: stars — stars: variables: other
1. Introduction
1.1. Previous observation of dust around O-rich stars
In the late sixties, while investigating deviations of stellar energy distributions from blackbodies, Gillett et al.
(1968) discovered a peak near 10 µm in the spectra of four
late-type, evolved, variable stars. Woolf & Ney (1969) attributed this emission peak to circumstellar silicate grains
around these stars. Since then there has been much interest in the exact nature of the dust around cool evolved
stars, how this dust forms and the structure of the dust
shells.
Hackwell (1972) suggested that the spectra of many
M-stars were not consistent with the view that the
circumstellar dust is comprised solely of silicate dust,
438
A.K. Speck et al.: IR spectra of O-rich evolved stars
implying that other constituents should be sought.
Treffers & Cohen (1974) made high spectral resolution
observations of oxygen-rich stars and concurred with
Woolf & Ney (1969) on the attribution of the circumstellar dust features to silicates, however they did not
preclude the inclusion of other grain types.
There have been various attempts to classify the different oxygen-rich dust features seen in the IRAS LRS spectra of evolved stars (e.g. Little-Marenin & Little 1990;
hereafter LML90, Sloan & Price 1995; hereafter SP95).
LML90 have classified the variation in the spectral features from M-type AGB stars into six categories: featureless, broad, 3 component, sil++ (a “9.7 µm” feature with
a strong feature on its long wavelength side centered at
about 11.3 µm), sil+ (a stronger “9.7 µm” feature with a
weaker long wavelength feature) and sil (a strong “9.7 µm”
silicate feature). They suggested that there is an evolutionary sequence in the spectral features, starting with a
featureless continuum, then developing a broad feature,
followed by a three component feature, a two component
feature and finally increasingly strong silicate features.
Following this work, SP95 devised another classification
system and found that they could categorize the spectra
into eight groups which formed a smooth progressive sequence from broad feature to “classic” silicate emission.
This classification system involved examining the fluxes
at 10, 11 and 12 µm (F10 , F11 and F12 ), and comparing
the ratios of the fluxes, rather than a visual inspection of
the feature shapes. Where their finding differed from those
of LML90 was with respect to the 12.5 − 13.0 µm feature.
They found that this feature was apparent in approximately 40−50% of the spectra across all groups. Therefore,
if the sequence is evolutionary, the 12.5 − 13.0 µm feature
is due to a dust-type whose formation is not related to
dust evolution.
Sloan et al. (1996) continued this work on the IRAS
LRS database and found that the 12.5 − 13.0 µm feature
was present primarily in the spectra of semiregular variables (SRs): 75 − 90% of SRb variables exhibit the feature,
while only 20 − 25% of Miras exhibit this feature. Other
than this strong tendency to be found in the spectra of
SRb variables, the 12.5 − 13.0 µm feature was not found
to be related to any peculiarities that would distinguish
the sources.
Onaka et al. (1989) had attributed the broad 12.5 µm
band seen in the spectra of some O-rich stars (see Hackwell
1972 and Vardya et al. 1986) to alumina (Al2 O3 ). They
also found that fits to nearly all their M star spectra could
be improved by the inclusion of Al2 O3 grains. They suggested that the only way to form the silicate grains was
for them to grow on pre-existing grains of Al2 O3 which
act as seed nuclei. However, for most polytypes of Al2 O3 ,
the infrared feature peaks at ∼11.5 − 12.0 µm rather than
12.5 − 13.0 µm. One polytype, α-Al2 O3 (or corundum),
has a peak in the right wavelength range, but also exhibits a feature at ∼21 µm which is not seen in the LRS
spectra. Therefore, Al2 O3 seems unlikely to be the cause
of the 12.5 − 13.0 µm band seen in these spectra. This
conclusion was supported by the work of Begemann et al.
(1997), who suggested that the 12.5 − 13.0 µm band is
probably associated with silicates.
Justtanont et al. (1998) observed strong emission lines,
attributed to CO2 , in the ISO-SWS spectra of stars which
exhibit the 12.5 − 13.0 µm feature. Their data showed
that the strengths of the CO2 emission lines and the
12.5 − 13.0 µm feature are well correlated, implying that
these features originate in the same region of the circumstellar envelope. Justtanont et al. (1998) asserted that the
occurrence of correlated CO2 and 12.5 − 13.0 µm emission is indicative of a warm (∼650 − 1250 K) gas layer
close to the star where the carriers of these features are
formed. They also suggested that an increase in mass-loss
would provide a mechanism to prevent formation and/or
excitation of the CO2 .
Posch et al. (1999) used ISO-SWS data to characterize the observed 12.5 − 13 µm feature and then compared
this to the spectra of various appropriate minerals. They
suggested that spinel (MgAl2 O4 ) is the most likely candidate. The attribution of the 12.5 − 13.0 µm feature to
this mineral is discussed in more detail in Sect. 3.
Although much the focus of recent research has been
on the nature of the 12.5 − 13.5 µm emission feature, a
number of other papers have focussed more on the general silicate features. Sloan & Price (1998; hereafter SP98)
extended their previous LRS dataset (SP95) to include
supergiants and S-stars. They found that the majority of
supergiants fell into the “classic” silicate groups defined
for AGB stars. There was little evidence for supergiants
with the broad feature. How their findings compare to the
current work will be discussed later.
Sylvester et al. (1994, 1998) discovered the presence
of 8.6– and 11.3 µm “UIR” emission bands (usually attributed to carbon-rich material, e.g. polycyclic aromatic
hydrocarbons) superposed on the silicate emission features
of a number of M supergiant stars, particularly amongst
members of the h&χPersei double cluster. Such UIR-band
emission was not found in the spectra of any AGB stars.
Sylvester et al. attributed this stark contrast in emission
properties to the presence of a chromospheric UV radiation field around the M supergiants, capable of dissociating the CO molecules which would otherwise have locked
up most of the available carbon atoms in the circumstellar
outflows. PAH molecules might then form from the liberated carbon atoms, and be fluorescently excited to emit in
the IR by the ambient UV photons. Sylvester et al. only
detected UIR-band emission from intermediate luminosity
supergiants which had relatively broad silicate bands. The
non-detectability of any UIR-band emission in their spectra of more luminous supergiants could be due to smothering of the UV photon field in the denser winds of these
stars, or to lack of contrast of the bands against the much
stronger silicate emission features of these stars.
A.K. Speck et al.: IR spectra of O-rich evolved stars
Sylvester (1999) investigated the 10 µm spectra of a
number of oxygen-rich evolved stars whose IRAS LRS
spectra had been classified as showing SiC 11 µm band
emission. Using UKIRT CGS3 spectra with higher S/N
and spectral resolution, he found that two of these stars,
both supergiants, instead showed 11.3 µm UIR-band
emission superposed on silicate features. The remainder
showed standard 3-component O-rich dust features, with
peaks at 10, 11 and 13 µm, apart from one object with a
self-absorbed silicate feature. There is therefore currently
no evidence for SiC emission features in the spectra of
any O-rich evolved stars.
1.2. Dust formation around oxygen-rich evolved stars
Scenarios for the formation of dust grains in circumstellar shells around oxygen-rich stars have been investigated
by several groups over the last thirty years. It seems appropriate to discuss the types of dust grains that these
models predict to appear in order to constrain our own
attributions.
The first attempts to predict the types of grains
that should be expected to form in these environments
neglected the actual nucleation processes and concentrated on the chemistry (e.g. Gilman 1969 and references
therein). The first investigation of dust formation processes in terms of nucleation and grain growth was by
Salpeter (1974) who concluded that grain formation
proceeded by nucleation of small refractory seed grains
(i.e. oxides) onto which an “onion-layer” mantle of the
more abundant silicates could form. Since then the basic
premise of grain nucleation and growth has remained
the same (e.g. Sedlmayr 1989; Jeong et al. 1999), but
the details change and the exact nature of the expected
condensates is still being debated.
Pégourié & Papoular (1985) discussed the grain-types
expected in circumstellar shells. Their model, based
mainly on models of condensation in the solar nebula,
implied that the precise nature of the grains formed is
determined by the elemental composition and oxidation
properties of the parent atmosphere, along with the
density structure of the dust shell. According to their
model: 1) the concentration of iron in silicates is always
expected to be low (mole % Fe2 SiO4 (fayalite) ∼ 20%;
and FeSiO3 ∼ 10%); 2) Mg2 SiO4 (forsterite) forms before
MgSiO3 (enstatite) in a cooling atmosphere, but forsterite
is converted into enstatite by gas-solid reaction with SiO.
They did not expect the dust to be pure forsterite or even
forsterite with some (≤ 20%) fayalite. Disequilibrium
calculations showed that the dust shells of M-stars should
also contain SiO2 , solid (metal) Fe, Ca2 SiO4 and Al2 O3 .
Stencel et al. (1990) hypothesized the formation of
“chaotic silicate”. In their scenario, a chaotic silicate forms
from a supersaturated vapour containing metal atoms,
SiO, AlO and OH in a hydrogen atmosphere. Inside the
439
chaotic silicate, where both silicon and aluminium are less
than fully oxidized, the higher reducing potential of Al
would initially act to produce Al-O bonds at the expense
of Si-O bonds. Thus, the stretching modes of solid, amorphous alumina would grow at the expense of the 9.7 µm
Si-O stretch associated with silicates. However, Si and O
are approximately ten times more abundant than Al and
therefore once the aluminium is completely oxidized, the
Si and SiO components of the grain should begin to oxidize and thus increase the strength of the classic 9.7 µm
silicate band.
Tielens (1990) reviewed the thermodynamic and kinetic factors which go into determining which grain species
form in O-rich circumstellar environments. This was taken
to be the “classic” condensation sequence and fits with the
basic ideas of grain nucleation and growth presented here.
This will be discussed in more detail in Sect. 4.
In a more recent paper, Kozasa & Sogawa (1997, 1998)
suggested that seed nuclei are alumina (Al2 O3 ) grains,
so that the first dust grains that form would be “naked”
Al2 O3 grains. These refractory grains can exist relatively
close to the star. Further out, where the temperature allows formation of lower temperature condensates, these
Al2 O3 grains would act as seed nuclei and become coated
with silicates. Such core-mantle grains would have a minimum size of ∼150 nm. This is very similar the Salpeter
(1974) model. However, Kozasa & Sogawa also suggested
that, even further from the star, homogeneous silicate
nucleation can occur, so that there would be a population of very small (a few nanometres) pure silicate grains.
They suggested that the classic 10 µm silicate feature
is due to these very small grains while the core-mantle
Al2 O3 -silicate grains are responsible for the 12.5 − 13 µm
feature. We will return to this question in Sect. 3.
The utilisation of Al2 O3 grains as seed nuclei was
called into question by Jeong et al. (1999) who suggested
that alumina is unlikely to be present in the gas phase and
therefore could not be used to build the “critical cluster”
that constitutes a seed nucleus. However, Gail & Sedlmayr
(1999) state that the relevant information needed to establish this is not available for alumina. Jeong et al.
(1999) used TiO2 as their nucleation seed material, after
finding that it was thermodynamically and chemically the
most favoured refractory species (see Gail & Sedlmayr
1998). Their models suggested that the first dust species
to form, closest to the star, would be aluminium silicates
(Al6 Si2 O13 ), a result not found in other models. They also
found that titanium-bearing dust should form. However,
Ti is over thirty times less abundant than Al and
400 times less abundant than Mg, Si, and Fe, which probably precludes the formation of detectable quantities of
Ti-rich dust. At somewhat lower temperatures, Jeong
et al. expected Mg2 SiO4 (forsterite), MgSiO3 (enstatite)
and SiO2 to form. At even lower temperatures they
predicted that iron oxides should form and eventually
contribute ∼30% of the dust volume, contrary to the
440
A.K. Speck et al.: IR spectra of O-rich evolved stars
Lodders & Fegley (1999) statement that “stars do not
rust”.
Gail & Sedlmayr (1999) predicted that M-star circumstellar shells should contain several different dust species.
Their models, like those of Jeong et al. (1999), start with
TiO2 cluster formation; however, as mentioned above,
Gail & Sedlmayr did not exclude the possibility of Al2 O3
seed nuclei. They found the most abundant condensate
to be olivine. The exact Fe/Mg ratio could vary and
they even suggested that this ratio changes as the grain
grows, so that there could be a gradient in the Fe/Mg ratio across the grains. They also found that metallic iron
particles form (in agreement with Lodders & Fegley), as
well as MgO (periclase) and a very small amount of SiO2
material.
It is clear from the models discussed above that there
is not a consensus as to the nature of the dust grains
that form around O-rich stars. However, the models do
provide insight into whether some dust species are at all
plausible. The aim of the present paper is to compare and
contrast the mid-infrared features found in the spectra
of cool evolved O-rich stars. To this end, mid-infrared
UKIRT CGS3 spectra have been obtained that have a
significantly higher signal to noise ratio and a somewhat
higher spectral resolution than those available in the IRAS
LRS database, although the 7.5 − 13.5 µm spectral coverage of these ground-based spectra is more restricted than
the 7.5 − 22 µm coverage of the IRAS LRS. We aimed in
particular to determine whether any dust feature types are
associated with particular types of star, and what can be
learned from the differences in spectral features amongst
the various stellar types.
2. The observational investigation
We present here a survey of dust features present in the
mid-IR spectra of 80 oxygen-rich AGB stars and 62 Msupergiants, with a view to identifying similarities and differences between the features amongst these various types
of stars. The selection of targets was largely dictated by
their accessibility to the 3.8-m UKIRT on Mauna Kea at
the dates of several different observing runs. The AGB
stars were observed mainly during two runs, in May 1991
and October 1992, although RW And and RZ Ari were
observed in October 1990 and SX Cyg was observed in
November 1993. The supergiant stars were observed over
a wider range of dates, from October 1992 to June 1998.
The supergiant spectra from October 1992 and August
1995 have already been published by Sylvester et al. (1994,
1998). They are included here because they are used for a
different purpose to that of Sylvester et al.
The nature of our survey means that these are snapshot spectra, obtained at a random phase of the stellar
pulsation cycle. Work by Little-Marenin et al. (1996),
Creech-Eakman et al. (1997) and Monnier et al. (1998)
Fig. 1. Mean spectra for each of the AGB star groups – progressing between the broadest feature (bottom right) and the
narrowest silicate feature (top left)
has indicated that AGB star silicate emission bands can
vary systematically as a function of pulsation cycle, primarily in the strength of the feature but sometimes also
in its profile. However, these observed profile variations
have in general been significantly less pronounced than
the differences between the emission feature types that
are classified here.
2.1. Observations
Observations were taken using the 3.8-m United Kingdom
Infrared Telescope (UKIRT) with the common-user spectrometer CGS3 (see Cohen & Davies 1995). We obtained
7.5 − 13.5 µm spectra with a 5.5-arcsec circular aperture,
and a spectral resolution of 0.17 µm. Two grating settings
gave a fully sampled 64-point spectrum. Wavelength calibration was with respect to observations of a Kr arc-lamp.
A.K. Speck et al.: IR spectra of O-rich evolved stars
441
Table 1. AGB stars observed
Source
IRAS names Other names
Variability
Type∗
IRC+60001, HD 225082
M
IRC-20007, AFGL 53, HD 1760
SRb
IRC+60009, AFGL 57, HD 1845
M
IRC+30015, AFGL 109, HD 4489, CIT 2
M
IRC+50049, AFGL 278, HD 11979
SR?
Mira, IRC+00030, AFGL 318, HD 14386
M
IRC+00032, AFGL 4195, HD 15105
M
IRC+50062, AFGL 335, HD 15186
M
IRC+20051
SRb
IRC-20043, AFGL 500, HD 22228
M
IRC+10050, AFGL 529, NML Tau
M
IRC-20049, AFGL 542, HD 25725
SRc?
IRC-30033, AFGL 552, HD 26601
M
IRC-10075, AFGL 615
SR
AFGL 617, HD 29844
M
IRC+60150, AFGL 664
M
IRC+50138, HD 33877
SR
IRC+50141, AFGL 715
M
IRC+00080, AFGL 786
M
IRC+40135, AFGL 794
M
IRC+40136, AFGL 802, HD 37645
M
IRC+30126, AFGL 822, HD 37724
M
IRC-20077, AFGL 8105
M
IRC-30049
M
AFGL 1008
M
IRC+00140,
M
AFGL 1028
M
IRC+50180, AFGL 1120, HD 58521
SR
IRC-20123, AFGL 1140, HD 60218
M
IRC-20134, AFGL 1151
M
IRC-10184, AFGL 1215, HD 65940
M
IRC+30209, AFGL 1329, HD 78712
SR
IRC+10215, AFGL 1380, HD 84748
M
IRC-20254, AFGL 1627, HD 117287
M
IRC-30207, AFGL 1650, HD 120285
SRa
IRC+30257, AFGL 1706, HD 126327
SRb
AFGL 4990S, HD 136753
M
IRC+30283, AFGL 1832, HD 145459, CIT 8
M
IRC+20298, AFGL 1858, HD 148206
M
IRC+20314,
M
M
IRC+10365, AFGL 2206
M
IRC+10360, AFGL 2213, HD 172171
M
M
IRC+10406, AFGL 2324, HD 177940
M
IRC-30403, HD177868
M
IRC-30404, AFGL 5556
M
IRC-10497, AFGL 2349
M
IRC-20555, HD 181060
M
M
IRC+30379, AFGL 2426
M
IRC+40355, AFGL 2429
M
IRC-30419, AFGL 5569, HD 188378
M
IRC-50314, HD 190163
M
IRC+30423, HD 192788
M
IRC-20590
M
IRC+10475
M
IRC-10546
M
IRC+00489
M
IRC+50347
SRa
IRC+60301
M
IRC-10498, AFGL 2775
M
IRC+50390, AFGL 2790, HD 206483
SRa
IRC+00509, AFGL 2806, HD 207076, SAO 145652
SRb
IRC+60337
M
IRC+40501, AFGL 2845, HD 209872
SRb
IRC+60345, AFGL 2865
SRb
IRC+10527, AFGL 3023, HD 218292
M
IRC+60389, HD 218997
M
IRC+40536, AFGL 3088
M
IRC-20642, AFGL 3136, HD 222800
M
IRC+60418, AFGL 3141, HD 222914
M
IRC+50484, AFGL 3188, HD 224490
M
Y Cas
T Cet
T Cas
RW And
SAO 37673
o Cet
R Cet
RR Per
RZ Ari
RT Eri
IK Tau
V Eri
W Eri
BX Eri
R Cae
TX Cam
UX Aur
R Aur
X Ori
RU Aur
SZ Aur
U Aur
RT Lep
S Col
CH Pup
AZ Mon
GX Mon
Y Lyn
Z Pup
DU Pup
U Pup
RS Cnc
R Leo
R Hya
W Hya
RX Boo
S CrB
RU Her
U Her
BG Her
V1692 Sgr
V1111 Oph
X Oph
V2059 Sgr
R Aql
AG Sgr
V342 Sgr
W Aql
Z Sgr
V635 Aql
BG Cyg
V462 Cyg
RR Sgr
Z Cyg
SX Cyg
RU Cap
Y Del
Y Aqr
W Aqr
RZ Cyg
UW Cep
UU Peg
RU Cyg
EP Aqr
YY Cep
SV Peg
CU Cep
R Peg
V Cas
BU And
R Aqr
Z Cas
R Cas
00007+5524
00192-2020
00205+5530
00445+3224
01556+4511
02168-0312
02234-0024
02251+5102
02530+1807
03318-1619
03507+1115
04020-1551
04094-2515
04382-1417
04387-3819
04566+5606
05121+4929
05132+5331
05351-0147
05367+3736
05384+3854
05388+3200
05404-2342
05450-3142
06434-3628
06551+0322
06500+0829
07245+4605
07304-2032
07329-2352
07585-1242
09076+3110
09448+1139
13269-2301
13462-2807
14219+2555
15193+3139
16081+2511
16235+1900
17072+1844
18320-1918
18349+1023
18359+0847
18501-2132
19039+0809
19044-2856
19093-3256
19126-0708
19167-2101
19343+0912
19369+2823
19384+4346
19528-2919
20000+4954
20135+3055
20296-2151
20392+1141
20417-0500
20438-0415
20502+4709
20581+5841
21286+1055
21389+5405
21439-0226
22000+5643
22035+3506
22097+5647
23041+1016
23095+5925
23212+3927
23412-1533
23420+5618
23558+5196
Spectral
Type
M7e
M5/M6I/II
M7e
S6,2e
M7
M7IIIe
M4e
M6e-M7e
M6III
M7e
M6e
M5/M6IV
M7e
M3
M6e
M9
M4(II)
M7
M8
M8e
M8e
M7e
M9e
M6e
Me
M6e
M9
M6S
M4-M9e
M
M5e-M8e
M6S
M8e
M7e
M8e
M7.5e
M7e
M7e
M7e
M3
M9
M9
M5e-M9e
M8
M7e
M5-M6e
M9
S6,6e
M5e
?
M7e
M7e
M5e
M5e
M7e
M9e
M8e
M6.5e
M7
M7
M8
M7e
M8e
M8
M6
M7
M5
M7e
M5.5e
M7e
M7e
M7e
M7e
Observation Calibrator
8.0 µm flux
Date
(10−12 Wm−2 µm−1 )
Oct. 5 ’92
Cyg
5.46
Nov. 1 ’93
βPeg
10.0
Oct. 5 ’92
Cyg
21.3
Oct. 5 ’90
αTau
1.97
Oct. 4 ’92
αCet
12.8
Oct. 6 ’92
αLyr
185
Oct. 4 ’92
αCMa
1.01
Oct. 7 ’92
αCet
3.20
Oct. 5 ’90
αTau
7.94
Oct. 4 ’92
αTau
5.42
Oct. 7 ’92
αCet
138
Oct. 4 ’92
αTau
7.48
Oct. 4 ’92
αCma
2.13
Oct. 4 ’92
αCMa
3.23
Oct. 4 ’92
αCMa
3.72
Oct. 4 ’92
αCma
14.3
Oct. 4 ’92
αCet
0.988
Oct. 4 ’92
αCet
13.2
Oct. 7 ’92
αAur
3.26
Oct. 7 ’92
αAur
3.99
Oct. 7 ’92
αAur
2.51
Oct. 7 ’92
αAur
6.33
Oct. 4 ’92
αCMa
2.05
Oct. 4 ’92
αCMa
1.30
Oct. 7 ’92
αCet
2.28
Oct. 7 ’92
αCet
1.26
Oct. 7 ’92
αCet
9.32
May 24 ’91
αBoo
5.40
Oct. 6 ’92
αCma
4.97
Oct. 6 ’92
αCMa
5.21
Oct. 6 ’92
αCma
1.88
May 25 ’91
αBoo
25.4
May 22 ’91
αBoo
67.4
May 25 ’91
αBoo
57.8
May 21 ’91
αBoo
152
May 24 ’91
αBoo
25.8
May 24 ’91
αBoo
7.08
May 24 ’91
αBoo
8.74
May 24 ’91
αBoo
12.6
May 24 ’91
βPeg
0.948
Oct. 7 ’92
βPeg
0.813
Oct. 4 ’92
αLyr
11.0
Oct. 4 ’92
αLyr
13.2
Oct. 4 ’92
αLyr
2.92
Oct. 6 ’92
αLyr
16.0
Oct. 7 ’92
βPeg
0.619
Oct. 5 ’92
βPeg
4.29
Oct. 7 ’92
βPeg
65.2
Oct. 5 ’92
βPeg
2.02
Oct. 7 ’92
βAnd
0.252
Oct. 4 ’92
αLyr
2.45
Oct. 4 ’92
αLyr
1.15
Oct. 7 ’92
βPeg
4.76
Oct. 5 ’92
βPeg
1.88
Nov. 2 ’93
αLyr
1.30
Oct. 7 ’92
βAnd
0.533
Oct. 7 ’92
βAnd
1.87
Oct. 7 ’92
βPeg
1.36
Oct. 7 ’92
βAnd
1.95
Oct. 5 ’92
βPeg
3.71
Oct. 5 ’92
βPeg
1.34
Oct. 7 ’92
Cyg
6.76
Oct. 5 ’92
βPeg
9.46
Oct. 6 ’92
αLyr
20.3
Oct. 5 ’92
βPeg
1.04
Oct. 7 ’92
Cyg
9.47
Oct. 5 ’92
βPeg
4.99
Oct. 5 ’92
βPeg
10.8
Oct. 5 ’92
βPeg
3.67
Oct. 6 ’92
βPeg
4.41
Oct. 6 ’92
αLyr
32.2
Oct. 7 ’92
Cyg
3.56
May 21 ’91
αBoo
41.5
∗
Variability Type: M = Mira; SR = semi-regular variable.
T Cet and V Eri are italicized because it is unclear whether each star is a giant or a supergiant – see Sect. 2.2.
442
A.K. Speck et al.: IR spectra of O-rich evolved stars
Table 2. Supergiant stars in the sample
∗
Source
IRAS names Other names
Variability
Type∗
IRC+60008,
Lc
IRC+50043, HD 10465
Lc?
IRC+60070, HD 236915
??
IRC+50052, HD 12041
SRc
IRC+60074, HD 13136
Lc
HD 13658
??
IRC+60078
SRc
IRC+60079, HD 14142
SRc
IRC+60081, HD 14242
??
IRC+60082, HD 14270
SRc
IRC+60083, HD 14330
Lc
IRC+60085, HD 14404
Lc
IRC+60086, HD 14469
SRc
IRC+60087, HD 14488
SRc
IRC+60088, HD 14528, AFGL 323
SRc
BD+56595
??
IRC+60090, HD 14826
??
IRC+60093, HD 236979
??
IRC+60094, AFGL 359
Lc
HD 237010, BD+57647
??
IRC+60110
Lc
IRC-30282, HD 155161
SRc
??
IRC-30326, HD 3167486, AFGL 2017
SRc
HD 163869
Lc
IRC-20431, AFGL 2071, HD 165674
SRc
??
IRC-10422, AFGL 2162
SRc
IRC-10435, AFGL 2186
??
IRC+00398
Lc
IRC+10420, AFGL 2390
??
??
IRC+20438, HD 339034
Lc
IRC+40409, AFGL 2560
Lc
IRC+40415, AFGL 2575
Lc
IRC+50351, AFGl 2683
??
IRC+60343, AFGL 2857
Lb
IRC+60357, AFGL 2916, HD 239978
Lc
IRC +50446, HD 215924, AFGL 2957
SRc
??
IRC+60410, AFGL 3110
Lc
MZ Cas
BD+47485
BD+58342
XX Per
KK Per
V550 Per
BU Per
T Per
V506 Cas
AD Per
FZ Per
PR Per
SU Per
RS Per
S Per
V439 Per
V441 Per
YZ Per
GP Cas
V648 Cas
IO Per
AH Sco
IRC-30312
KW Sgr
V540 Sgr
VX Sgr
IRC-10419
UY Sct
HD 171094
UW Aql
V1302 Aql
IRC-20565
NR Vul
BC Cyg
KY Cyg
AZ Cyg
AZ Cep
ST Cep
U Lac
V355 Cep
V358 Cas
00186+5940
01400+4815
01550+5901
01597+5459
02068+5619
02116+5754
02153+5711
02157+5843
02167+5926
02169+5645
02174+5655
02181+5738
02185+5622
02188+5652
02192+5821
02196+5658
02217+5712
02347+5649
02360+5922
02473+5738
03030+5532
17080-3215
17374-3156
17488-2800
17566-3555
18050-2213
18227-1347
18248-1229
18305-1408
18550+0023
19244+1115
19272-1929
19480+2447
20197+3722
20241+3811
(20h56m)
22069+5918
22282+5644
22456+5453
22471+5902
23281+5742
Spectral
Type
M1.3Iab
M2Ib
M2.4Ib
M3.6Ib
M1.9Ib
M5.4Iab
M3.7Ib
M2.1Iab
M5.7Iab
M2.4Iab
M0.3Iab
M0.7Iab
M3.3Ib
M4.4Ib
M4.5Iab
M5.8Iab
M3.1Iab
M1.9Iab
M2.8Iab
M2.9Iab
M3I
M5I
M2.6Ia
M2.4Ia
M5Iab
M5-M6I
M2.5Iab
M3.4Iab
M3I
M2.2Iab
F8I
M2Ib
M1.1Iab
M3.2Iab
M3.9Iab
M3.1Iab
M1.6I
M2.6Ia
M2.5Ia
M1.1Iab
M2.8Ia
Observation Calibrator
8 µm flux in
Date
10−12 Wm−2 µm−1
Oct. 5 ’92
Cyg
1.25
Aug. 16 ’95
β And
0.660
Aug. 16 ’95
β And
0.433
Oct. 6 ’92
β Peg
2.58
Oct. 7 ’92
α Cet
0.889
Oct. 7 ’92
α Cet
0.274
Oct. 4 ’92
α Tau
0.705
Oct. 4 ’92
α Tau
0.400
Oct. 7 ’92
α Cet
0.591
Oct. 4 ’92
α Tau
0.771
Oct. 7 ’92
α Cet
0.521
Oct. 7 ’92
α Cet
0.552
Oct. 4 ’92
α Tau
1.12
Oct. 6 ’92
β Peg
1.33
Oct. 4 ’92
α Tau
6.81
Oct. 7 ’92
α Cet
0.368
Oct. 7 ’92
α Cet
0.929
Oct. 7 ’92
α Cet
0.976
Aug. 16 ’95 β And
0.983
Aug. 16 ’95
β And
1.01
Oct. 4 ’92
α Tau
4.88
Aug. 16 ’95
η Sgr
14.0
Aug. 16 ’95
η Sgr
1.68
Aug. 16 ’95
η Sgr
3.50
Aug. 16 ’95
η Sgr
1.10
Oct. 4 ’92
βPeg
73.3
Aug. 17 ’95
γ Aql
0.450
Jun. 27 ’98
γ Aql
3.72
Aug. 17 ’95
γ Aql
0.875
Jun. 27 ’98
γ Aql
0.945
Oct. 6 ’92
α Lyr
18.2
Aug. 17 ’95
γ Aql
0.285
Aug. 16 ’95
γ Aql
2.46
Aug. 16 ’95
γ Aql
7.47
Aug. 16 ’95
γ Aql
6.95
Aug. 16 ’95
γ Aql
2.26
Aug. 16 ’95
β Peg
0.480
Aug. 17 ’95
β Peg
1.01
Aug. 17 ’95
β Peg
2.95
Aug. 16 ’95
β Peg
0.442
Aug. 16 ’95 β And
1.42
Variability Type: L = irregular variable; SR = semi-regular variable.
Table 3. 10 µm classifications of AGB stars
featureless AGB
BU And
R Hya
R Peg
RZ Ari
T Cas
UX Aur
V Cas
broad AGB
AG Sgr
BG Cyg
R Aql
R Aur
R Leo
RR Sgr
RT Eri
S Col
SZ Aur
V1692 Sgr
W Aql
W Hya
YY Cep
V Eri
X Oph
broad+sil AGB
BX Eri
RR Per
RW And
RX Boo
RZ Cyg
SAO 37673
UW Cep
V462 Cyg
W Aqr
Y Cas
Z Cas
SV Peg
AZ Mon
silicate AGB A
CU Cep
RS Cnc
RU Cyg
U Her
UU Peg
X Ori
Mira
silicate AGB B
DU Pup
GX Mon
R Cet
RU Cap
S CrB
U Aur
V1111 Oph
V342 Sgr
V635 Aql
Z Cyg
Z Pup
silicate AGB C
IK Tau
R Aqr
RT Lep
RU Aur
RU Her
EP Aqr
TX Cam
U Pup
V2059 Sgr
Y Del
Y Lyn
Z Sgr
silicate AGB D
CH Pup
R Cae
R Cas
W Eri
Y Aqr
A.K. Speck et al.: IR spectra of O-rich evolved stars
443
Fig. 2. Continuum-subtracted AGB star spectra classed as showing broad features. x-axis is wavelength in µm. y-axis is flux in
W m−2 µm−1
444
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 3. Continuum-subtracted AGB star spectra classed as showing broad+sil features. x-axis is wavelength in µm. y-axis is flux
in W m−2 µm−1
A.K. Speck et al.: IR spectra of O-rich evolved stars
445
Table 4. 10 µm classifications of supergiant stars
broad Super A
AD Per
FZ Per
V506 Per
PR Per
V441 Per
T Per
KK Per
BD+47485
broad Super B
BD+58342
V439 Per
V550 Per
silicate Super 1
AH Sco
KW Sgr
silicate Super 2
AZ Cyg
BC Cyg
BU Per
HD 171094
IRC-20565
ST Cep
V358 Cas
V540 Sgr
XX Per
YZ Per
silicate Super 3
AZ Cep
V648 Cas
GP Cas
IRC-10419
KY Cyg
MZ Cas
NR Vul
RS Per
SU Per
UW Aql
V355 Cep
silicate Super 4
V1302 Aql
IRC-30312
S Per
UY Sct
VX Sgr
Fig. 4. Continuum-subtracted AGB star spectra classed as showing group A silicate features. x-axis is wavelength in µm. y-axis
is flux in W m−2 µm−1
446
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 5. Continuum-subtracted AGB star spectra classed as showing group B silicate features. x-axis is wavelength in µm. y-axis
is flux in W m−2 µm−1
The telescope secondary was chopped east-west at 5 Hz
using a 30-arcsec throw. The error bars in the plots represent 1σ standard errors on the fluxes. Due to imperfect
cancelation of the time-varying atmospheric ozone band,
the error bars are larger at approximately 9.6 µm. Details
of the observed stars can be found in Table 1 for the AGB
stars and Table 2 for the supergiants.
2.2. Classifying the features
For each source in our sample a 3000 K blackbody representing the stellar photosphere was normalized to the
spectrum at 8.0 µm. This was then subtracted from the
observed astronomical spectrum to yield the spectra that
are plotted in Figs. 2–8 and 10–14. Since the mean level of
the spectra between 7.5 and 8.0 µm is zero and the error
bars are larger because of greater atmospheric absorption
A.K. Speck et al.: IR spectra of O-rich evolved stars
447
Fig. 6. Continuum-subtracted AGB star spectra with group C silicate features. x-axis is wavelength in µm. y-axis is flux in
W m−2 µm−1
between these wavelengths, we have only plotted our
continuum-subtracted spectra from 8.0 µm onwards. Our
continuum subtraction procedure implicitly assumes that
the 8.0 µm continuum is dominated by blackbody-like stellar photospheric emission, i.e. that neither continuum dust
emission nor gaseous SiO fundamental-band emission or
absorption is important. Higher resolution spectra of the
SiO fundamental and first overtone bands (e.g. Tsuji et al.
1997; Waters et al. 1999; Aringer et al. 1999) indicate that
the gaseous SiO band is not likely to have significantly
perturbed our low resolution spectra. A much more detailed modeling of wide spectral coverage ISO-SWS spectra would be needed to investigate whether continuum
dust emission in general makes a significant contribution
at 8 µm. However, the ISO-SWS spectra of Tsuji et al.
(1997) indicate that for high mass loss stars, such as the
448
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 7. Continuum-subtracted AGB star spectra classed as showing group D silicate features. x-axis is wavelength in µm. y-axis
is flux in W m−2 µm−1
scatter that is expected for these stars, should not affect
our data analysis.
Fig. 8. The continuum-subtracted spectrum of T Cet. x-axis is
wavelength in µm. y-axis is flux in W m−2 µm−1
supergiant S Per, continuum dust emission is significant at
8 µm but that for lower mass loss rate stars (the majority
of our sample) the contribution from continuum dust emission at this wavelength is much less pronounced.
In producing our continuum-subtracted spectra (CSS),
we have assumed that the photospheric temperature for
each star is 3000 K. However, we find that varying the
adopted blackbody temperature by ±1000 K about the
adopted 3000 K has very little effect on the shape and
strength of the derived continuum-subtracted dust features, since we are in the Rayleigh-Jeans domain of the
photospheric Planck distributions. Therefore the exact
effective temperatures of the stars, within the natural
In order to aid comparisons between different spectra, we also produced normalized “continuum-subtracted
spectra” (CSS) for the AGB stars and supergiants in our
sample, by normalizing the CSS to a peak flux of unity.
The normalized CSS of each AGB star was compared to
those of the other AGB stars to find similarities and form
groups. Leaving aside the seven AGB stars in our sample whose individual spectra show no pronounced emission features (classed as “featureless” in Table 3), the
dust features of the AGB stars have been classified into
three groups: broad AGB, where the feature extends from
8 µm to about 12.5 µm with little structure; broad+sil
AGB, which consists of a broad feature with an emerging 9.7 µm silicate bump; and silicate AGB, which is the
“classic” 9.7 µm silicate feature. In addition, the silicate
AGB group can be classified further into four subgroups,
A, B, C, D, which show slightly differently shaped silicate
features, as shown by the mean spectra. These variations
seem to show a sequence from broader to narrower silicate feature. The stars in each group are summarized in
Table 3, and the mean feature profile for each group is
shown in Fig. 1. The original CSS are shown in Figs. 2−7,
grouped in the classes described above.
The broad+sil AGB groups can also be split into
two groups according to whether the spectra show the
12.5 − 13.0 µm feature. This will be discussed later.
Two of our sources, T Cet and V Eri, have been
classified both as AGB stars and as supergiants in various
A.K. Speck et al.: IR spectra of O-rich evolved stars
449
published works. For this reason, exactly whether these
stars should be included in the AGB star groupings
or the supergiant groupings is ambiguous. In fact
T Cet has been classified as an AGB star (e.g.
Bedding & Zijlstra citeB98), a supergiant (e.g.
SP98), M-star (e.g. Bedding & Zijlstra 1998), MS-star
(e.g. Groenewegen & de Jong 1998) and S-star (e.g.
Skinner et al. 1990; Groenewegen 1993). This source also
has a markedly different spectrum from any others in the
sample (e.g. Speck 1998; Skinner et al. 1990). We can
see in Fig. 8 that the spectral features for this source are
unlike any others in our sample and it therefore needs to
be treated separately. For this reason, T Cet was removed
from the sample. V Eri has variously been classed as
an AGB star and as a supergiant (see e.g. Hashimoto
1994; Habing 1996; Neri et al. 1998; Triglio et al. 1998;
Cernicharo 1998 etc.). Kwok et al. (1997) give V Eri the
spectral classification M6II, i.e. intermediate between
a giant and a supergiant, while Houk & Smith-Moore
(1988) assign V Eri a spectral class of M5/M6IV, which
would make it a subgiant. Furthermore, V Eri exhibits the
12.5 − 13.0 µm feature, which is seen in many AGB star
spectra but only in one other supergiant spectrum (that
of S Per; see Sect. 3). We have therefore placed it in the
AGB grouping, rather than in the supergiant grouping.
In the same way as for the AGB stars, the normalized CSS of each supergiant was compared to those of the
other supergiants to find similarities and form groups. The
dust features of the supergiant stars can also be classified
into three basic groups: featureless, whose emission above
the continuum is weak and hard to classify; broad Super,
where the feature extends from ∼9 µm to ∼13 µm; and
silicate Super, which again is the “classic” 9.7 µm silicate
feature. Since the featureless spectra show little of interest regarding dust they will not be discussed further here.
The broad feature can be classified into two further subgroups, one extending from 9 to 13 µm, and one with a
short wavelength onset at 9.5 µm. The difference between
these two broad groups may be entirely due to differences
in the underlying dust continuum. Examination of the silicate features in these spectra is slightly compromised by
the appearance of UIR bands (see Sylvester et al. 1994,
1998). The strong 11.3 µm UIR band has been edited out
of the averaged spectra shown in Fig. 9. Again, the “classic” silicate group (silicate Super) can be classified further
into four subgroups, 1, 2, 3 & 4, where the exact shapes
of the silicate features vary slightly and the averages show
this. One supergiant with a silicate feature, U Lac, did not
fit into any of these four groups, but does have the basic
silicate feature and matches group B from the AGB star
classifications. The stars in each group (excluding featureless) are summarized in Table 4, the mean feature from
each group is shown in Fig. 9, and the original CSS are
shown in Figs. 10−14, grouped in the classes described
above.
Fig. 9. Mean spectra for each of the supergiant star groups progressing between the broadest feature (bottom right) and
the narrowest silicate feature (top left)
Our classifications of AGB stars and supergiant can be
compared to those of SP95 and SP98. These comparisons
are shown in Tables 5 and 6. The classification system used
by SP95 & SP98 gives each spectrum a number: SE1 to
SE8, where SE1 refers to the broad feature; the SE number increases as the silicate feature emerges up to SE8
which refers to the “classic” narrow silicate feature. We
would therefore expect spectra in the our broad feature
class to have low SE numbers. As we progress through the
classes up to the narrow silicate feature class, the SE number should increase. In general we concur with the SP95
and SP98 progression, as can be seen in Tables 5 and 6.
However, there are a couple of spectra for which the classifications are markedly different: W Hya, which we class
as having a broad feature while SP95 give it a strong silicate feature; and R Cae, which we class as silicate D and
they give a broad feature class. These differences may
be due to their classification method, which only uses
450
A.K. Speck et al.: IR spectra of O-rich evolved stars
relative fluxes at 10, 11 and 12 µm rather than the overall
shape of the features. However, in comparing the classifications used we found that none of the stars that we have
classified as broad Super appear in their dataset and they
do not have any examples of the broad feature. This is
unfortunate, since it appears to be one of the few ways
in which the observed dust features differ between AGB
stars and supergiants.
2.3. Comparing dust features from AGB stars and
supergiants
Having classified the spectra into groups, we have compared the spectral features for AGB stars and supergiants.
2.3.1. The “classic” silicate feature
Amongst both the AGB and supergiant classes we have
distinguished sub-groupings of the “classic” silicate feature. Comparing these sub-groups, we find that the various silicate features found in the spectra of AGB stars
are very similar to those found in the spectra of supergiants. Figures 15a−d shows the comparisons of the various silicate features. In each case the mean AGB silicate feature for each sub-group is plotted along with the
closest matching supergiant silicate features. Figure 15a
compares AGB silicate group A with two supergiant silicate groups (1 & 2). The AGB feature is found to match
the average of these two supergiant features. Figure 15b
shows AGB silicate group B, together with the spectrum
of the unclassified supergiant U Lac. Figure 15c shows
AGB silicate group C with two supergiant silicate groups
(3 & 4). As with AGB silicate group A, this AGB group
appears to be intermediate between the two supergiant
groups. Figure 15d shows AGB silicate group D with the
supergiant group 4. It is clear from these figures that, although the silicate feature varies from star to star, similar shaped features are present in the spectra of both
AGB stars and supergiants, and seem to form a progression between broader and “classic” narrow 9.7 µm silicate
features.
2.3.2. The broad feature
As we can see from Fig. 16, the broad AGB star feature
is quite different from the two broad supergiant features.
The broad features for supergiants are slightly narrower
than for AGB stars, extending from 9.0 − 9.5 µm to
∼13 µm, rather than 8 − 12.5 µm as in the AGB star spectra. Furthermore, the flux continues to drop longward of
12.5 µm, rather than leveling off as in the case of the AGB
stars. In addition, the broad feature for AGB stars seems
to develop a silicate peak on its way to becoming a “classic” silicate feature (see broad+sil classification; Fig. 1);
this does not happen in the supergiant spectra. Only one
supergiant in the sample, S Per (see Fig. 14), exhibits the
12.5 − 13.0 µm feature seen in some AGB star spectra
and this may be due to this star’s unusual characteristics.
It is believed that S Per is a relatively new supergiant and
that its dust shell is thinner and more spherically symmetric than is usually found for supergiants (Richards
1997; Richards et al. 1999). These characteristics make
the dust shell more like that of a semiregular red giant
than a supergiant and may explain the appearance of the
12.5 − 13.0 µm feature in its spectrum.
2.4. The featureless spectra
Those AGB and supergiant 8 − 13 µm excess spectra
which have been labeled as “featureless” are all very weak
compared to the local stellar photospheric continua. Signal
to noise considerations may therefore be largely responsible for the inability to classify these excess spectra. To
test this, we co-added the five AGB star 8 − 13 µm excess
spectra labeled as “featureless” in Table 5. The resulting
mean spectrum had better S/N and resembled the broad
AGB star feature discussed above. It may therefore be the
case that the respective broad features are responsible for
the low-contrast “featureless” spectra of AGB stars and
supergiants.
3. The ∼13 µm feature in AGB star spectra
As discussed above, there is virtually no evidence for
the 12.5 − 13.0 µm feature in supergiant spectra. It
is, however, clearly present in some AGB star spectra
(see Figs. 2−7; see also the ISO-SWS observations of
Posch et al. 1999) and requires further investigation. The
12.5 − 13 µm feature is not ubiquitous amongst the AGB
star spectra. In our sample, this feature is found rarely in
the spectra which exhibit a broad feature extending from
8.5 to 12 µm, sometimes in spectra with the “classic” silicate feature, and is most regularly found in the spectra
that combine the broad and silicate features, although it is
still not universal. The AGB star spectra in the broad+sil
group, where the 12.5 − 13.0 µm feature is most frequently
seen, have been divided further into those spectra which
exhibit the 12.5 − 13.0 µm feature, and those which do
not. The mean features of these two subgroups are plotted together in Fig. 17a, where the 12.5 − 13.0 µm feature is clearly seen. In Fig. 17b, we have subtracted the
mean of the broad+sil spectra which do not exhibit the
13.0 µm feature from the mean of those spectra that do,
in order to isolate this feature. From this we have found
the peak position of the mean 12.5 − 13 µm feature to
be 12.92 µm. This is very close to the position measured
by Posch et al. (1999), 12.95 − 13.05 µm, although is notably on the short wavelength side of their measurement.
A.K. Speck et al.: IR spectra of O-rich evolved stars
451
Fig. 10. Continuum-subtracted supergiant spectra classed as showing broad features. x-axis is wavelength in µm. y-axis is flux
in W m−2 µm−1
This discrepancy may be due to the feature’s proximity to
the long wavelength limit of our observed spectra, which
may have caused the loss of a small portion on the long
wavelength side of the feature. This would also account
for the smaller full width half maximum (FWHM) of our
mean 12.5 − 13 µm feature of 0.45 µm, compared to the
FWHM for this feature of 0.5−0.8 µm measured by Posch
et al. (1999) in their higher resolution SWS spectra spectra of 11 stars. Their mean spectrum shows the feature
to have a pronounced long-wavelength asymmetry, which
could however be largely due to CO2 and other narrow
emission features which are present in this wavelength region of their spectra.
If we look for differences between the stars which exhibit the 13 µm feature, and those which do not, we
find that, within the broad+sil group, all the stars with
the feature are non-Mira semi-regular variables, whereas
those stars without the feature are all Miras (see Table 7).
452
A.K. Speck et al.: IR spectra of O-rich evolved stars
Table 5. Comparing classifications - AGB star features
Source
BU And
R Hya
R Peg
T Cas
V Cas
BG Cyg
R Aql
R Leo
RR Sgr
RT Eri
S Col
SZ Aur
W Aql
W Hya
YY Cep
V Eri
X Oph
RR Per
RW And
RX Boo
RZ Cyg
UW Cep
V462 Cyg
W Aqr
Y Cas
Z Cas
SV Peg
AZ Mon
Variability
Type
M
M
M
M
M
M
M
M
M
M
M
M
M
SR
M
SR
M
M
M
SR
SR
M
M
M
M
M
SR
M
Our
Class
featureless
featureless
featureless
featureless
featureless
broad
broad
broad
broad
broad
broad
broad
broad
broad
broad
broad
broad
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
broad+sil
SP95
Class
SE2
SE2t
SE4
SE1
SE1
SE2t
SE5
SE2
SE4
SE3
SE2
SE2t
SE3
SE8
SE4
SE3t
SE1
SE2
SE3
SE3t
SE3t
SE3
SE1
SE3
SE3
SE1
SE3t
SE3t
There are three sources (AZ Mon, UW Cep and V462 Cyg)
which show slight perturbations at 12.5 − 13 µm, however it is unclear whether these sources exhibit the feature or not. Therefore these sources have been classified
as not showing the feature. Furthermore, in the other spectral feature groups, where the 12.5 − 13.0 µm feature is
less common, we find the same situation: those stars with
the feature are non-Mira semi-regular variables and those
stars without the feature are Miras. As discussed earlier,
Sloan et al. (1996) found that 75−90% of the SR variables
in their sample exhibited the feature, while it was seen in
only 20 − 25% of their Mira spectra. This is similar to the
results of Begemann et al. (1997), who found that 30% of
the Miras in their sample exhibited the 12.5 − 13.0 µm feature. The smaller size of our sample may explain why we
do not find any 12.5 − 13.0 µm feature amongst the Mira
spectra. However, the result still stands that the major difference between stars with or without the 12.5 − 13.0 µm
feature is the variability type. We therefore need to ask
the question: why is this feature more visible in SR spectra than in Mira spectra? Is it that there is a dust species
that forms around SRs but not around Miras? Or is it that
Source
CH Pup
R Cae
R Cas
W Eri
Y Aqr
IK Tau
R Aqr
RT Lep
RU Aur
RU Her
EP Aqr
TX Cam
U Pup
Y Lyn
Z Sgr
DU Pup
GX Mon
R Cet
RU Cap
S CrB
U Aur
V1111 Oph
V342 Sgr
Z Cyg
Z Pup
CU Cep
RU Cyg
U Her
UU Peg
X Ori
o Cet
Variability
Type
M
M
M
M
M
M
M
M
M
M
SR
M
M
SR
M
M
M
M
M
M
M
M
M
M
M
SR
SR
M
M
M
M
Our
Class
silicate D
silicate D
silicate D
silicate D
silicate D
silicate C
silicate C
silicate C
silicate C
silicate C
silicate C
silicate C
silicate C
silicate C
silicate C
silicate B
silicate B
silicate B
silicate B
silicate B
silicate B
silicate B
silicate B
silicate B
silicate B
silicate A
silicate A
silicate A
silicate A
silicate A
silicate A
SP95
Class
SE4
SE3
SE5
SE6
SE4
SE5
SE6
SE6
SE8
SE5
SE4
SE5
SE6
SE8
SE6
SE5
SE6
SE6
SE8
SE5
SE4
SE5
SE6
SE8
SE6
SE6
SE5
SE4
SE5
SE4
SE8
we are seeing high temperature refractory grains in SRs
which become coated with other material around Miras?
How would this explain the appearance of this feature in
the spectrum of the supergiant S Per?
3.1. Miras vs. Semiregular variables
Since we have found a difference in the dust features for
Miras and semiregular variables, it seems appropriate to
discuss the differences between types of red variables. The
basic categories are: Miras, which have regular pulsation
with large amplitude (V ≥ 2.5 mag) and periods longer
than 60 days; semiregular (SR) variables, which have less
regular pulsation and smaller amplitudes (V ≤ 2.5 mag);
and irregular (L) variables. Furthermore, the SRs can be
split into four subgroups; SRa, which have relatively regular periods; SRb with more irregular periods; SRc, which
are more luminous and associated with supergiants; and
SRd which are warmer as defined by their colours or spectral types. However, the differences and similarities between these groups and their relationships to one another
A.K. Speck et al.: IR spectra of O-rich evolved stars
453
Table 6. Comparing classifications - Supergiant silicate features
Source
AH Sco
KW Sgr
BU Per
ST Cep
V358 Cas
V540 Sgr
XX Per
KY Cyg
NR Vul
RS Per
SU Per
UW Aql
S Per
UY Sct
Variability
Type
SRc
SRc
SRc
Lc
Lc
Lc
SRc
Lc
Lc
SRc
SRc
Lc
SRc
SRc
Our
Class
Sil1
Sil1
Sil2
Sil2
Sil2
Sil2
Sil2
Sil3
Sil3
Sil3
Sil3
Sil3
Sil4
Sil4
SP98
Class
SE5t
SE6
SE7
SE7
SE6
SE6
SE7
SE5
SE4
SE5t
SE4
SE5
SE4
SE4
Table 7. AGB stars in the broad+silicate group with and
without the ∼13 µm feature
with
∼13 µm
BX Eri
RX Boo
RZ Cyg
SAO 37673
SV Peg
without
∼13 µm
AZ Mon∗
UW Cep∗
V462 Cyg∗
RR Per
RW And
UW Cep
W Aqr
Y Cas
Z Cas
variable
type
SR
SR
SR
SR
SR
Mira
Mira
Mira
Mira
Mira
Mira
Mira
Mira
Mira
∗
These spectra may exhibit the feature very weakly.
However it is so weak that we feel it is best to class these
spectra as not showing the 13 µm feature.
have caused some disagreement in recent years and therefore require further discussion.
It has been suggested in the past that SRs were less
evolved than Miras (e.g. Feast & Whitelock 1987), however, given that apart from the 12.5 − 13.0 µm feature,
the observed mid-IR spectral features are basically identical for the SR variables and Miras and cover the full range
of classifications, this is apparently not reflected in the
state of evolution of the dust. Kerschbaum & Hron (1992)
also suggested that the SRs and Miras form a continuous sequence, with the SRs the progenitors of the Miras.
Furthermore, it has recently been suggested that evolution from SR to Mira is unlikely to be monotonic and
that stars may alternate between being Miras and SRs
(P. Whitelock, pers. comm.; Kerschbaum & Hron 1992).
The Miras and SRs have also been found to have separate
Fig. 11. Continuum-subtracted supergiant spectra classed as
showing group 1 silicate features. x-axis is wavelength in µm.
y-axis is flux in W m−2 µm−1
period-luminosity relations (Whitelock 1986), but whether
the separation of these two relations is due to differences
in pulsation or stellar structure is unknown (Bedding &
Zijlstra 1998). Habing (1996) has suggested that a combination of dust shell structure and mass-loss rate evidence
for SRs and Miras implies that Miras are in fact the progenitors of the SR variables. Add this to the uncertainties
raised over whether SRs are indeed a separate class to the
Miras (Kerschbaum & Hron 1992), and we have a very
confused picture.
Mattei et al. (1997) re-investigated the differences between Miras and semiregular variables and found that
the groups were indeed distinct from one another based
on an examination of the amplitudes of variability and
their stabilities. Miras have much more stable amplitudes,
while the amplitudes of SRs are much more variable.
Furthermore, they were able to divide the SRs into two
subgroups based on the amplitude of pulsation and the
multiperiodicity. These groups are basically SRa and SRb.
Thus, we have evidence that the red variables can be
classified into discrete groups rather than a continuous
sequence, although whether one variability type is the
progenitor of another is still unclear.
Hron et al. (1997) examined dust features in the
IRAS-LRS spectra of Mira and SR variables and found
a difference between these two stellar types. While they
suggested that the Miras are just an extension of the
SRs with lower effective temperature, they also found
that both Miras and SRs exhibit the full range of mid-IR
emission features, from broad to the classic, narrow 9.7 µm
feature. Furthermore, they found an intriguing difference
between these variability types. The Mira spectra behave
as we would expect if the sequence from broad to narrow
silicate feature is indeed evolutionary, with the optical
depth increasing as the feature narrows (i.e. more dust
= more evolved dust). However, the SRs showed no such
trend and in fact, based on radiative transfer modeling,
tended to have more optically thin dust shells for the
narrower features (Ivezic & Elitzur 1995). Furthermore,
Ivezic & Knapp (1998) found that there was a link
454
A.K. Speck et al.: IR spectra of O-rich evolved stars
between mass-loss rate and variability type for AGB stars.
They found that, while Miras could be fit by “standard”
steady-state radiatively driven models, with inner dustenvelope radii defined by the typical condensation temperature (800 − 1000 K), their models required SRs to have
a much lower condensation temperature of ∼300 − 400 K.
They interpret this as a decrease in the mass-loss, which
equates to a less dense dust shell, in agreement with
Hron et al. (1997). This is evidence that the SRs should
be treated as a distinct group from the Miras, at least as
far as dust formation is concerned.
As mentioned above, it has also been suggested that
the transition between Miras and SRs is not a single event
and happens many times (e.g. Kerschbaum & Hron 1992).
A star may have regular, strong, large amplitude pulsations for a time and then have weaker, more irregular pulsation and then go back to strong, regular pulsation as
the thermal pulses of the star disrupt the stellar structure. Therefore, it is possible that the 12.5 − 13.0 µm
feature needs to be explained in terms of a dust species
that can only form in the environment of an SR, and is
either rapidly destroyed or over-coated by a lower condensation temperature material in the environment of the
Miras or has been ejected by the stellar wind prior to the
Mira phase. The small percentage of Miras found by Sloan
et al. (1996) to exhibit the 12.5 − 13.0 µm feature might
be explained in terms of stars that have recently changed
from SR to Mira and have not yet destroyed/lost their
complement of the carrier of the 12.5 − 13.0 µm feature.
As discussed in the introduction, Justtanont et al.
(1998) found a correlation between CO2 emission lines
and the 12.5 − 13.0 µm feature. This was taken as
evidence that both the CO2 and the carrier of the
12.5 − 13.0 µm feature were formed in warm gas layers
close to stars with low mass-loss rates, in particular
the lower density dust shells around SRs. Since the
strengths of the 12.5 − 13.0 µm and CO2 emissions are
correlated and higher mass-loss rates appear to preclude
formation/excitation of the CO2 lines, higher mass-loss
rates may also inhibit the formation of the 12.5 − 13.0 µm
carrier.
3.2. Which minerals?
Previously, the 12.5 − 13.0 µm feature has been identified with Al2 O3 grains (Hackwell 1972; Vardya et al.
1986; Onaka et al. 1989; LML90, Glaccum 1995), however there are problems with this attribution. Al2 O3 , or
alumina is found in several different forms. Only one form
occurs naturally on the earth: corundum or α-Al2 O3 . This
has rhombohedral crystal structure and is known by various other names including ruby or sapphire depending
on the trace impurities (Deer et al. 1966, hereafter DHZ).
There are known synthetic polytypes: β-Al2 O3 , which has
hexagonal crystal structure; and γ-Al2 O3 which is cubic.
However both these forms convert to corundum on heating (DHZ). It is also possible to make amorphous samples
(e.g. Begemann et al. 1997). Some condensation models
(Salpeter 1974; Sedlmayr 1990; Tielens 1990; Kozasa &
Sogawa 1997, 1998) suggest that Al2 O3 should be the first
dust type to condense around O-rich stars. According to
some previous research (e.g. Vardya et al. 1986; Onaka
et al. 1989; Tielens 1990; Kozasa & Sogawa 1997, 1998),
Al2 O3 then forms a nucleation seed on which the silicates
can form a mantle. Alternatively, Nuth & Hecht (1990)
and Stencel et al. (1990) suggested that the condensate is
a “chaotic silicate” in which, initially, the emission from
Al-O bonds dominates the spectra, but is then overwhelmed by emission from the more abundant Si-O bonds.
Both sets of authors agreed that the 12.5 − 13.0 µm band
seen in the spectra of oxygen-rich stars can be attributed
to Al-O bonds and that it signifies the presence of some
form of aluminium oxide. Moreover, Al2 O3 grains found in
meteorites (Nittler et al. 1994a,b, 1997; Huss et al. 1995)
have isotopic signatures which suggest they were formed
around oxygen-rich AGB stars. However, the abundance
of such AGB star Al2 O3 grains is very low (< 1 ppm; cf.
6 ppm for presolar SiC and 400 ppm for presolar diamond).
Furthermore, the polytype (i.e. the crystal structure) of
these presolar Al2 O3 grains is as yet unknown (L. Nittler,
pers. comm.) so that no information on the exact spectral
features we should expect has yet been gleaned from these
meteoritic grains.
Begemann et al. (1997) studied the laboratory spectra
of various forms of aluminium oxide, both crystalline and
amorphous, with a view to identifying more clearly the
12.5 − 13.0 µm feature seen in IRAS-LRS spectra. They
found that amorphous aluminium oxide could not account
for the observations and that, while the α- crystalline form
of Al2 O3 could account for the 12.5 − 13.0 µm feature,
a second feature seen at 21 µm in laboratory spectra of
this material was not observed in astronomical spectra.
Indeed, the two spectra of Al2 O3 shown in Fig. 18 have
peak positions in the 11.5 − 12.0 µm range, rather than at
12.5 − 13.0 µm. Begemann et al. (1997) suggested that the
12.5 − 13.0 µm feature may come from a form of silicate
rather than from aluminium oxide. Furthermore, we may
note that if the 12.5 − 13.0 µm feature was attributable to
the first condensate, it would be expected to appear predominantly in the broad feature spectra. As stated above,
the 12.5 − 13.0 µm feature appears most commonly in our
spectra in the broad+sil class, and is certainly not ubiquitous. In addition, the appearance of the feature seems to
be related to the variability type of the AGB star (i.e. Mira
or SR). Given that Al2 O3 is so important to the theoretical condensation sequence, and that the spectral features
of Al2 O3 do not seem to match the observed astronomical
features, it seems doubly unlikely that the 12.5 − 13.0 µm
feature is attributable to Al2 O3 .
Kozasa & Sogawa (1997, 1998) proposed grains consisting of Al2 O3 cores and silicate mantles for the carrier
A.K. Speck et al.: IR spectra of O-rich evolved stars
455
Fig. 12. Continuum-subtracted supergiant spectra classed as showing group 2 silicate features. x-axis is wavelength in µm. y-axis
is flux in W m−2 µm−1
of the 12.5 − 13.0 µm feature. However, this was investigated by Posch et al. (1999) who found that: 1) there
would have to be a large population of grains of ∼85%
Al2 O3 ; and 2) there would have to be an even larger number of pure silicate particles to produce anywhere near the
correct ratio of 13 µm to 10 µm flux strengths. They, therefore, concluded that such core-mantle grains were unlikely
to be the carriers of the 12.5 − 13.0 µm feature.
Another possible carrier for the 12.5 − 13.0 µm
feature is spinel, MgAl2 O4 , as proposed by Posch et al.
(1999) who compared laboratory and calculated spectra
of various candidate minerals with ISO-SWS spectra of
AGB stars which exhibit the 13 µm feature. Posch et al.
argue against Al2 O3 as the carrier for many of the reasons
summarized here, and concluded that the most likely mineral to produce the 13.0 µm feature is spinel (MgAl2 O4 ).
456
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 13. Continuum-subtracted supergiant spectra classed as showing group 3 silicate features. x-axis is wavelength in µm. y-axis
is flux in W m−2 µm−1
Spinel is a very refractory, cubic mineral which is found
naturally on the earth (DHZ). It has also been found as a
presolar grain in meteorites (Nittler et al. 1994a, 1994b,
1997). However, this attribution is compromised by problems with the spinel optical data. Posch et al. (1999) used
optical data from Tropf & Thomas (1991) which is potentially flawed, since the optical constants are based on a
compilation of different data mostly from synthetic
samples. Tropf & Thomas noted that where more than
one data set exists for the same wavelength region,
the measurements do not necessarily match up. Following
Andreozzi et al. (2000), it has become clear that, for
spinel in particular, the precise level of order/disorder of
the crystal structure can have a large effect on the optical
A.K. Speck et al.: IR spectra of O-rich evolved stars
457
Fig. 14. Continuum-subtracted supergiant spectra classed as showing group 4 silicate features. x-axis is wavelength in µm. y-axis
is flux in W m−2 µm−1
properties. The Mg/Al order varies among the synthetic
samples and affects the IR spectra in both the width and
number of peaks. Thus, the optical constants n and k
depend on Mg/Al order, which in turn may depend on
the formation temperature. This renders compilations of
data for spinel, such as that of Tropf & Thomas (1991),
moot since such compilations do not take into account
the issue of the differing amounts of order/disorder in the
differing samples. Therefore the positions of the spectral
features derived by Posch et al. (1999) from the Tropf &
Thomas data are unlikely to be valid for any one sample
of spinel. Previous studies of crystalline spinel have
found peaks at wavelengths longwards of 13.5 µm (e.g.
Chopelas & Hofmeister 1991; Hafner 1961; Preudhomme
& Tarte 1971). Indeed, it was these previously published
spectra which prompted Speck (1998) to ignore spinel in
her discussion of mid-IR dust features because it did not
appear to have any features in the 7.5 − 13.5 µm window
and would therefore not be a candidate for the 13 µm
feature. Clearly the issue of spinel and its intriguing IR
spectral behaviour requires more investigation, but that
is beyond the scope of the current work.
Our results support the findings of Begemann et al.
(1997) by associating the 13 µm feature with silicates.
This observed feature appears to be best explained in
terms of some form of silicate (Begemann et al. 1997) or
SiO2 (silica; Speck 1998). Both copper (Cu(II)SiO4 ) and
zinc (ZnSiO4 ) silicates exhibit this spectral feature (Speck
1998), but it is not seen in the spectra of the more commonly expected magnesium, calcium or aluminium pyroxene or olivine silicates. Copper and zinc are not abundant
enough to form observable quantities of such silicates. The
silicates of magnesium, calcium and aluminium which are
expected to form in these circumstellar environments are
pyroxenes (i.e. chain silicates) and olivines (i.e. individual
units of silicate tetrahedra; see Speck 1998 for a discussion
of silicate structures). Silicon dioxide, on the other hand,
forms a “fully polymerized” crystal lattice. Intermediate
between SiO2 and the pyroxenes are the sheet silicates.
From a spectral feature point of view, there is a progression from SiO2 , with a relatively strong 13 µm feature
(Fig. 18d), through the sheet silicates, where the 13 µm
feature diminishes, to pyroxenes which generally do not
show this feature. Although silica is chemically an oxide,
the structures and properties of the silica minerals are
more closely allied to those of silicates. SiO2 can be either crystalline or amorphous. There are three main polytypes for crystalline SiO2 : quartz, tridymite and cristobalite. These forms may be seen as a progression in the
temperature of formation (DHZ). Silica can also exist in
amorphous forms such as lechaterlierite, obsidian and silica glass. The different forms of silica and their stable
temperatures are discussed in more detail by Speck (1998).
Furthermore, the available atomic abundances are perfectly acceptable for the formation of observable quantities
of SiO2 . Therefore, it is plausible on abundance grounds
that silicon dioxide or a sheet silicate is responsible for
the 13 µm feature, so an SiO2 origin for the 13 µm feature
deserves further exploration.
4. Dust condensation sequence
The dust condensation sequence in O-rich circumstellar
regions is expected to go as follows (Tielens 1990):
Al2 O3 , Alumina
⇓
Ca2 Al2 SiO7 , Augite
⇓
CaMgSi2 O6 , Diopside
⇓
MgAl2 O4 , Spinel
⇓
CaAl2 Si2 O8 , Anorthite
Mg2 SiO4 , Forsterite
⇓
MgSiO3 , Enstatite
⇓
(Mg, Fe)2 SiO4 , Olivine
According to condensation thermodynamics, the silicate
condensation sequence starts with the nucleation of alumina (Al2 O3 ) from the circumstellar gas at about 1760 K.
The first silicate is expected to form by a gas-solid reaction with alumina, to form Ca2 Al2 SiO7 . As the temperature drops, further gas-dust reactions occur so that Mg
substitutes for Al to form CaMgSi2 O6 . The aluminium released and the remaining alumina are converted to spinel
(MgAl2 O4 ). As further cooling occurs the CaMgSi2 O6
458
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 16. Comparison of the average broad features for AGB stars
(solid line) and supergiants (dotted and dashed lines)
Fig. 15. Comparison of the silicate features of AGB stars and
supergiants. a) AGB silicate group A (solid line) vs. supergiant
silicate groups 1 (dotted line) and 2 (dashed line); b) AGB silicate group B (solid line) vs. unclassified supergiant U Lac
(dashed line); c) AGB silicate group C vs. supergiant silicate
groups 3 (dashed line) and 4 (dotted line); d) AGB silicate
group D vs. supergiant silicate group 4 (dashed line)
and the spinel form a solid-solid reaction, producing the
feldspar anorthite (CaAl2 Si2 O8 ). At even lower temperatures (∼ 1440 K) forsterite (Mg2 SiO4 ) starts to condense
out. Forsterite continues to form until the temperature
has dropped to ∼ 1350 K when it reacts with gaseous SiO
to produce enstatite (MgSiO3 ). Finally, at ∼ 1100 K reactions with gaseous iron will convert some enstatite into
fayalite (Fe2 SiO4 ) and forsterite. Kinetics also plays an
important role in determining which silicates form in the
outflows of AGB stars. Depending on the density structure of the circumstellar region the condensation sequence
will cease at different points. Thus, if the density drops
rapidly with distance from the star, the only dust expected to form will be various high temperature oxides
(e.g. Al2 O3 , CaTiO3 , ZrO2 ), which will form very close
to the star. If the densities are a little higher further
out in the circumstellar shell gas-grain reactions can take
place, allowing the formation of calcium-aluminium silicates. If the density is high enough a little further out
still magnesium silicates may form as rims on the Ca-Al
silicates. For magnesium silicates to nucleate, there need
to be very high densities a long way out, which is highly
unlikely. Feldspars, such as anorthite (CaAl2 Si2 O8 ) are
not expected to form, as the solid-solid reaction requires
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 17. The 12.5 − 13.0 µm feature. a) the mean spectra for
stars in the broad+silicate group, with (solid line) and without (dashed line) the ∼13 µm feature; b) the isolated ∼13 µm
feature. This was achieved by subtracting the mean spectrum
without the ∼13 µm feature from the mean spectrum with the
∼13 µm feature. The dashed line shows the peak position of
the feature. FWHM = full width at half maximum
unrealistically high densities. Finally, Fe can only be incorporated into Mg-silicates if, initially, most of the iron
is in gaseous form (rather than solid, metal form) and if
the density is high enough at large distances from the star
where fayalite (Fe2 SiO4 ) can survive. Figure 18 shows the
laboratory spectra of major minerals that are expected to
form, together with that of SiO2 which is a likely to be a
step in the formation of silicates.
459
Fig. 18. The laboratory spectra of some minerals in the
predicted condensation sequence. In a) two spectra are
shown for amorphous alumina, where the solid line represents a porous sample and the dashed line is a compact
sample. These spectra both come from the Jena optical constants database (http://www.astro.uni-jena.de/
Group/Subgroups/Labor/Labor/odata.html), as do the spectra in b) and c). In d), two polytypes of SiO2 are shown.
The solid line is a form of amorphous silica (from Nyquist
1971) and the dashed line is one of the crystalline polytypes,
tridymite (Hofmeister et al. 1992). The x-axis is wavelength
in µm, the y-axis is normalized extinction
460
A.K. Speck et al.: IR spectra of O-rich evolved stars
Another possible sequence of condensation mentioned
by SP98 involves chemistry determined by the C/O ratio.
In this case, less evolved stars with low C/O ratios would
exhibit silicate features while somewhat more evolved
stars, with C/O ratios closer to unity, could have spectra
dominated by Al2 O3 . According to the standard scenario
(e.g. Salpeter 1974), stars with low C/O ratios (1), have
an abundance of oxygen atoms to form dust and therefore
we should see strong silicate features. However, in stars
with higher C/O ratios (but still less than unity) less oxygen would be available and therefore aluminium oxide features could dominate because aluminium is more oxidizing
than magnesium and so uses up the available oxygen first
(see SP98 for details).
5. Discussion
The first obvious manifestation of the differences between
AGB and supergiant spectra is the difference in their broad
features. However, the various “silicate” features are basically identical for both AGB stars and supergiants (see
Fig. 15). Therefore, assuming that the observed spectral
features represent evolution of the dust from the most
refractory species, responsible for the broad features, to
magnesium-rich silicates, it is clear that the AGB star
condensation sequence is different from that of supergiants but both condensation sequences eventually leading
to similar dust types, i.e. magnesium-rich silicates. The
suggested evolutionary paths for the dust around O-rich
AGB stars and supergiants are shown in Table 8. To explain these we need to have distinct differences between
the dust-forming regions of AGB stars and supergiants in
terms of elemental abundances, oxidation properties, temperatures or densities of the dust-forming atmospheres.
The difference between the broad features exhibited by O-rich AGB stars and supergiants implies a
different pathway to the magnesium-rich silicates. The
short-wavelength rise of the supergiant broad feature
may be indicative of calcium-aluminium-rich silicates,
e.g. Ca2 Al2 SiO7 , which match this part of the spectrum
well (Fig. 19a). However, thus far no Ca-Al-rich silicate
spectrum has a broad enough feature to account for the
entire broad band. A second spectral feature peaking
at about 11.5 − 12.0 µm and extending to beyond the
13.5 µm limit of our spectra is needed to account for the
long wavelength part of the broad feature. This feature
(but not the 12.5 − 13.0 µm feature) could be attributable
to alumina (Al2 O3 ). Likewise the short wavelength onset
of the AGB star broad feature may be due to Mg-Fe-rich
silicates, but again this does not account for the long
wavelength part of the feature. However, the shape of the
long wavelength feature needed to complete the broad
feature in both the AGB star and supergiant cases is
very similar. We have found that this extra feature is
best fit using alumina (Al2 O3 ), although the exact form
Table 8. The different evolutionary paths of dust around AGB
stars and around supergiants
AGB stars
Supergiants
(Mg, Fe)2 SiO4 & Al2 O3
⇓
(Mg, Fe)2 SiO4
strengthens
Ca2 Al2 SiO7 & Al2 O3
⇓
No obvious
transition
feature
⇓
(Mg, Fe)2 SiO4
overwhelms
broad feature
⇓
(Mg, Fe)2 SiO4
overwhelms
broad feature
13 µm
feature
in some spectra
No 13 µm
feature
of alumina is not identical in each case to account for the
flux longwards of ∼ 13 µm. Figure 19a shows that the
supergiant broad feature can be fit using a combination of
Ca-Al-rich silicate (Ca2 Al2 SiO7 ) and compact amorphous
alumina, while the AGB star broad feature (Fig. 19b) can
be fitted using a combination of olivine (MgFeSiO4 ) and
porous amorphous alumina.
Therefore the main difference between the supergiant
and AGB dust features seems to be that the supergiants have a distinct evolutionary phase in which calciumaluminium-rich silicate condensation takes place to an extent that is observable, whereas for AGB stars this phase
is not seen. The first distinct spectral features to be seen
in AGB star spectra can be attributed to a combination
of Mg-Fe-rich silicates and alumina. The AGB stars go
through a phase in which the Mg-Fe-rich silicate feature
strengthens and eventually overwhelms the alumina feature. The supergiant spectra exhibit either Ca-Al silicate
and alumina features or strong Mg-Fe silicate features. No
transitional states have been observed.
The 12.5 − 13.0 µm feature seen in the spectra of some
AGB stars is almost completely absent from the supergiant spectra, appearing only in the spectrum of S Per,
a supergiant with a particularly thin dust shell similar
to that of SRs (see Sect. 2.3.2; Richards 1997; Richards
et al. 1999). Therefore, whatever dust species is responsible for this feature is not included in the usual condensation sequence for supergiants. As discussed in Sect. 3,
the 12.5 − 13.0 µm feature is predominantly seen in the
spectra of semiregular variables, so there must also be a
difference between dust-forming properties of Miras and
semiregular variables. We suggest that this feature is not
due to Al2 O3 , MgAl2 O4 , or core-mantle alumina-silicate
grains, as has previously been proposed (e.g. Vardya et al.
1986; Posch et al. 1999; Kozasa & Sogawa 1997, 1998
A.K. Speck et al.: IR spectra of O-rich evolved stars
Fig. 19. Fits to broad features. a) The supergiant broad feature
is fit by Ca2 Al2 SiO7 (dotted line) and compact amorphous
Al2 O3 (dashed line). b) The AGB star broad feature is fit by
MgFeSiO4 (dotted line) and porous amorphous Al2 O3 (dashed
line)
respectively), but rather that it may be associated with
silicon dioxide (see Morioka et al. 1998; Speck 1998), which
also accounts for the 9.25 µm peak seen in the spectra of
some AGB stars with strong silicate features (see Speck
1998). Another possible carrier may by highly polymerized silicates which bear more structural resemblance to
SiO2 than to isolated silicates like olivine.
The difference between the chemistries of Miras,
semiregular variables and supergiants needs to be addressed. There are two proposed dust sequences (for
want of a better term): the classic condensation sequence
discussed in Sect. 4 (see also Tielens 1990); and the new
chemical evolution sequence based on C/O ratios (SP98).
The SP98 model was constructed using a dataset that
had no supergiant spectra exhibiting the broad feature.
With this supergiant broad feature included, we need to
re-examine the possibilities.
AGB stars experience the so-called “third dredge-up”
which brings the products of the helium-burning shell to
461
the surface of the star (see e.g. Iben & Renzini 1981). In
particular, this process enriches the atmospheres of these
stars with 12 C. With successive dredge-ups the atmospheres of AGB stars get progressively more carbon-rich.
Therefore oxygen-rich AGB stars can have carbon-tooxygen ratios (C/O) ranging from cosmic, ∼0.4, to approximately unity. Supergiants, on the other hand, are not
expected to experience a “third dredge-up” and are therefore expected to remain oxygen-rich, with approximately
cosmic C/O. Furthermore, Sylvester et al. (1994) found
UIR bands in the spectra of some M-supergiants. They
suggested that UV photodissociation of CO molecules provided the free carbon atoms required to form the organic
materials responsible for the UIR bands. Obviously, this
dissociation mechanism would also release extra oxygen
atoms, and if the carbon atoms become trapped in organic molecules, this could lead to an even more oxygenrich environment (C/O < 0.4). Therefore, supergiants
exhibiting the UIR bands could have even more oxygen available for dust formation and, following the SP98
scheme, we might expect these stars to exhibit the “classic” strong silicate spectral feature, since they should have
more than enough oxygen to make substantial amounts of
Mg-silicates. However, many of the supergiants with UIR
bands exhibit “broad” spectral features. These supergiant
“broad” spectral features are alone enough to call into
question a condensation mechanism based on the amount
of available oxygen, since the C/O ratio for supergiants is
always low. Combining this with the UIR band evidence,
it seems unlikely that the type of dust forming around
these stars is determined solely by the amount of oxygen
available.
Contrary to SP98, we have found that the “broad” feature in AGB star spectra cannot solely be attributed to
Al2 O3 . The feature appears to be due to a combination of
Al2 O3 and “classic” silicate (e.g. amorphous olivine). A fit
to the average “broad” feature spectrum using Al2 O3 and
amorphous olivine is shown in Fig. 19b. This does not contradict the SP98 dust evolution scheme, but implies that
oxygen availability is never so low as to preclude some silicate formation. It does, however, cause a problem for the
supergiant sequence. The SP98 scheme cannot account for
the formation of Ca-Al-rich silicates rather than Mg-Ferich silicates in the very oxygen-rich supergiant environments, where there should be more than enough oxygen
to form the “classic” silicates.
It is clear that the C/O ratio is an important factor in determining which dust species will form around
which stars, however, other factors (e.g. dust shell temperature and density) are also important and need to be
taken into account. In trying to accommodate as much of
the available information, and previous interpretations, as
possible, the current work suggests the following: the type
of dust species formed around a given star is controlled
by the C/O ratio and by the density and temperature of
the dust-forming region, but that one of these factors will
462
A.K. Speck et al.: IR spectra of O-rich evolved stars
dominate, depending on the exact balance between the
C/O ratio, the density and the temperature. The work
of Hron et al. (1997) suggested that the classical condensation sequence is valid for Miras, since they found that
the “classic” narrow silicate feature is associated with the
most optically thick dust shells (i.e. more dust = more
evolved dust). This is not the case for the semiregular variables. In fact, for these the optical depth is lowest for the
narrowest silicate features and is higher for the broader
features (Ivezic & Elitzur 1995). However, the most optically thick SR dust shells are only as optically thick as the
most optically thin Mira dust shells. This may be the key
to the problem. In the case of the Miras, with relatively
high optical depths, and therefore relatively high density
dust shells, it is the density of the dust shells that controls
the dust formation, and these environments undergo dust
formation which follows the classic condensation sequence.
For Miras the evolutionary stage of the dust reflects the
evolutionary stage of the star (i.e. Miras with broad features are less evolved than those with the “classic” silicate
features). In the case of SRs, with much less dense shells,
it is the C/O ratio that controls the dust formation and
therefore the dust “evolution” follows the opposite path to
the classic condensation sequence, as suggested by SP98.
This fits with the findings of Hron et al. (1997). In the case
of the SRs, the stars with the “classic” silicate feature are
less evolved, with less dense dust and more free oxygen,
while the “broad” feature is associated with more evolved
SRs, which have denser dust shells and fewer available
oxygen atoms.
We now return to the enigmatic 13 µm feature. This
feature appears to be associated with SR variables and
therefore with less dense dust-forming regions. Following
the above argument, this in turn implies that species formation is dominated by the C/O ratio. However, it is not
clear that the 13 µm feature is associated with any other
dust feature. While Begemann et al. (1997) found a close
correlation between this feature and the classic silicate feature, both Sloan et al. (1996) and the current work found
that the 13 µm feature is not particularly associated with
any of the other dust features – appearing in examples
from all classes and implying that the 13 µm feature is not
associated with the evolutionary stages of the dust around
SRs, or the prevailing C/O ratio. Furthermore, it is rarely
found in supergiant or Mira spectra. It has been suggested
that the 13 µm feature is carried by silicon dioxide (SiO2 ;
Speck 1998). However, some objections to this attribution
have been raised. Firstly, there is the exact position of
the feature at ∼12.9 µm. Posch et al. (1999) state that
the feature in amorphous SiO2 peaks at 12.3 µm, which
is too far removed from the observed feature. However,
infrared spectra of different polytypes of SiO2 presented
by Speck (1998) show that the “13 µm” feature peak position varies from one polytype to another from ∼12.5 µm
to ∼12.8 µm, with the feature width (FWHM) varying
from 0.33 to 0.73 µm. While this is still not identical to
the observed feature, it is much closer than any other suggested carrier. Second, there are other spectral features
in the infrared spectrum of SiO2 which need to be accounted for. As shown in Fig. 18, SiO2 has two features
in the 7 − 14 µm range relevant to our observed spectra, the 12.5 − 13 µm feature and a stronger 9.2 µm feature. There is also a feature at ∼20 µm outside the range
of our observations. The relative strength of the 13 µm
feature to the other features depends on the polytype
(i.e. crystal structure; see Speck 1998), the level of disorder of the crystal structure, and the relative optical depths
(Hofmeister et al. 2000). While the 13 µm feature is usually weaker than the 9.2 µm feature, optical depth effects
can hugely increase the relative strength of the 13 µm
feature (see Hofmeister et al. 2000). The observed 13 µm
feature is never found in isolation − there is always some
contribution in the 8 − 10 µm region, probably from
Si-O bond stretching in amorphous forsterite/enstatite.
It is possible to bury the 9.2 µm feature in this general
Si-O stretch region, so that the only evidence for SiO2
above the general silicate contributions to the spectra is
the 13 µm feature. In fact, there is some evidence for an
emerging 9.2 µm feature, over and above the “classic”
9.7 µm silicate feature, which may be attributable to SiO2
(Speck 1998). It is possible that the ∼20 µm SiO2 feature
could be buried in the amorphous silicate Si-O bending
feature nearby, but this cannot be assessed with the current observations. However, it was noted by LML90 and
Goebel et al. (1989) that the 13 µm feature seen in IRAS
spectra was often associated with an unidentified feature
at about 19 − 20 µm, and this may be the other SiO2
spectral feature. Another objection raised to SiO2 is that
equilibrium calculations have shown that it is not expected
to form around AGB stars (Lodders & Fegley 1999). Any
SiO2 that does form is expected to be rapidly converted
to forsterite and then enstatite. The key point, however,
is that objection requires that equilibrium is reached. In
order to form forsterite and enstatite, the condensation
sequence is likely to go through SiO2 formation. Current
thinking has been that SiO2 forms at most only a minor
constituent of the dust, but this is based on the absence
of observed features near 9.2 and 12.5 µm and a shoulder near 8.4 µm (J. Nuth, pers. comm), which, as stated
above, may be buried in the strong Si-O stretching feature.
If equilibrium is not reached, we would expect there to be
some residual SiO2 in these regions. In Miras, the denser
dust shells are more likely to reach equilibrium than the
less dense SR dust shells because each atom/molecule in
a denser environment has more chance of interacting with
another atom/molecule. This might explain the formation
of a disequilibrium product (SiO2 ) in the less dense SR
dust shells and not in the denser Mira dust shells.
In Sect. 3, we discussed the difference between SRs
and Miras. It is not clear whether SRs are the progenitors
of Miras or vice versa, whether these stars go through
successive SR and Mira phases, or whether they are
A.K. Speck et al.: IR spectra of O-rich evolved stars
completely separate objects. However, assuming that they
are related and that SRs may evolve into Miras and then
back into SRs, how can we account for the appearance of a
species in the dust-forming region of one of these variability types and not in the other? The easiest explanation is
to assume that the dust around SRs, which we assume includes some SiO2 or highly polymerized silicate, is ejected
by the stellar wind prior to the Mira phase and then does
not form again until the star returns to the SR phase.
Alternatively, the dense dust regions around Miras may
cause the SiO2 -rich (parts of) grains to be transformed
into olivines and pyroxene as the progression towards equilibrium is enhanced in the denser environment. It is known
that the implantation of impurities, such as metal ions,
tends to transform SiO2 and highly polymerized silicates
into less polymerized silicates like olivines and pyroxenes
(Zachariasen 1932; Hess 1995). The less dense dust forming regions around SRs would allow more of the purer
SiO2 /highly polymerized silicate grains to form and survive. All other features in the spectra of Miras and SRs
are identical. The fact that SiO2 /highly polymerized silicate is plausible from an abundance point of view and that
the disappearance of the species in Mira spectra can be
explained makes the attribution of the 13 µm feature to
SiO2 or highly polymerized silicates even more attractive.
While laboratory spectra of SiO2 are commonly available,
as discussed above, the position of the 12.5 − 13.0 µm feature is not ideal. Therefore, laboratory studies of highly
polymerized silicates are needed to find a better match to
the observed feature.
How do we relate these ideas to the supergiants? Most
supergiant spectra do not exhibit the 13 µm feature. They
seem to follow a different chemical evolution to either
Miras or SRs. While many supergiant spectra show “classic” silicate features identical to those in AGB stars, it
is the difference in the broad feature that is the mystery. We have to ask: what differences between supergiants and AGB stars can shed light on this problem?
One obvious difference is that some supergiants exhibit
UIR emission bands, attributed by Sylvester et al. (1994,
1998) to the presence of chromospheres around such supergiants, whereas AGB stars do not have chromospheres.
The presence of a strong UV radiation field in supergiant
outflows, as well as dissociating CO and permitting the
formation of the carbon-rich UIR-band carriers, may also
significantly influence the early dust formation sequence in
the remaining oxygen-rich gas once the UIR-band carriers
have formed.
Another avenue for future investigation is the relation
of asymmetry factors for AGB stars to their dust spectra. It has been suggested that the asymmetry of the light
curve of an AGB star is related to its evolutionary stage
(Vardya et al. 1986). Unfortunately it appears that useful published data is only available for the light curves of
Miras, so that comparison of the light curves of Miras to
those of SRs is not currently possible. Other factors to take
463
into account in future work are dust feature strength and
profile variations as a function of pulsation cycle phase,
as studied by Monnier et al. (1998) and earlier workers,
and the secular long-term variation of the observed dust
features of some sources that was discovered by Monnier
et al. (1999).
Acknowledgements. We would like to thank Icko Iben,
Teije de Jong, Katharina Lodders, Larry Nittler, Joe Nuth,
Hans Olofsson, Anita Richards, Patricia Whitelock, and
Elric Whittington for useful email correspondence/chats.
An anonymous reviewer is also thanked for constructive criticism, which was instrumental in significantly improving this
paper. AKS would like to dedicate this work to the memory of
J. Colin Siddons, physics teacher, mentor and friend, who died
in November 1999.
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